A population of Main Belt Asteroids co-orbiting with Ceres and Vesta

We have carried out a search for Main Belt Asteroids (MBAs) co-orbiting with the large MBA Vesta and the dwarf planet Ceres. Through improving the search criteria used in Christou (Christou, A.A. [2000b]. Astron. Astrophys. 356, L71–L74) and numerical integrations of candidate coorbitals, we have identified approximately 51 (44) objects currently in co-orbital libration with Ceres (Vesta). We show that these form part of a larger population of transient coorbitals; 129 (94) MBAs undergo episodes of co-orbital libration with Ceres (Vesta) within a 2 Myr interval centred on the present. The lifetime in the resonance is typically a few times ～10⁵ yr but can exceed 2 × 10⁶ yr. The variational properties of the orbits of several co-orbitals were examined. It was found that their present states with respect to the secondary are well determined but knowledge of it is lost typically after ～2 × 10⁵ yr. Objects initially deeper into the coorbital region maintain their coorbital state for longer. Using the model of Namouni et al. (Namouni, F., Christou, A.A., Murray, C.D. [1999]. Phys. Rev. Lett. 83, 2506–2509) we show that their dynamics are similar to those of temporary coorbital NEAs of the Earth and Venus. As in that case, the lifetime of resonant libration is dictated by planetary secular perturbations, the inherent chaoticity of the orbits and close encounters with massive objects other than the secondary. In particular we present evidence that, while in the coorbital state, close encounters with the secondary are generally avoided and that Ceres affects the stability of tadpole librators of Vesta. Finally we demonstrate the existence of Quasi-Satellite orbiters of both Ceres and Vesta and conclude that decametre-sized objects detected in the vicinity of Vesta by the DAWN mission may, in fact, belong to this dynamical class rather than be bona-fide (i.e. Keplerian) satellites of Vesta.     

Delivery of dark material to Vesta via carbonaceous chondritic impacts

NASA’s Dawn spacecraft observations of Asteroid (4) Vesta reveal a surface with the highest albedo and color variation of any asteroid we have observed so far. Terrains rich in low albedo dark material (DM) have been identified using Dawn Framing Camera (FC) 0.75 μm filter images in several geologic settings: associated with impact craters (in the ejecta blanket material and/or on the crater walls and rims); as flow-like deposits or rays commonly associated with topographic highs; and as dark spots (likely secondary impacts) nearby impact craters. This DM could be a relic of ancient volcanic activity or exogenic in origin. We report that the majority of the spectra of DM are similar to carbonaceous chondrite meteorites mixed with materials indigenous to Vesta. Using high-resolution seven color images we compared DM color properties (albedo, band depth) with laboratory measurements of possible analog materials. Band depth and albedo of DM are identical to those of carbonaceous chondrite xenolith-rich howardite Mt. Pratt (PRA) 04401. Laboratory mixtures of Murchison CM2 carbonaceous chondrite and basaltic eucrite Millbillillie also show band depth and albedo affinity to DM. Modeling of carbonaceous chondrite abundance in DM (1–6 vol.%) is consistent with howardite meteorites. We find no evidence for large-scale volcanism (exposed dikes/pyroclastic falls) as the source of DM. Our modeling efforts using impact crater scaling laws and numerical models of ejecta reaccretion suggest the delivery and emplacement of this DM on Vesta during the formation of the ～400 km Veneneia basin by a low-velocity (<2 km/s) carbonaceous impactor. This discovery is important because it strengthens the long-held idea that primitive bodies are the source of carbon and probably volatiles in the early Solar System.     

Photometric,spectral phase and temperature effects on 4 Vesta and HED meteorites: Implications for the Dawn mission

Phase angle and temperature are two important parameters that affect the photometric and spectral behavior of planetary surfaces in telescopic and spacecraft data. We have derived photometric and spectral phase functions for the Asteroid 4 Vesta, the first target of the Dawn mission, using ground-based telescopes operating at visible and near-infrared wavelengths (0.4–2.5 μm). Photometric lightcurve observations of Vesta were conducted on 15 nights at a phase angle range of 3.8–25.7° using duplicates of the seven narrowband Dawn Framing Camera filters (0.4–1.0 μm). Rotationally resolved visible (0.4–0.7 μm) and near-IR spectral observations (0.7–2.5 μm) were obtained on four nights over a similar phase angle range. Our Vesta photometric observations suggest the phase slope is between 0.019 and 0.029 mag/deg. The G parameter ranges from 0.22 to 0.37 consistent with previous results (e.g., Lagerkvist, C.-I., Magnusson, P., Williams, I.P., Buontempo, M.E., Argyle, R.W., Morrison, L.V. [1992]. Astron. Astrophys. Suppl. Ser. 94, 43–71; Piironen, J., Magnusson, P., Lagerkvist, C.-I., Williams, I.P., Buontempo, M.E., Morrison, L.V. [1997]. Astron. Astrophys. Suppl. Ser. 121, 489–497; Hasegawa, S. et al. [2009]. Lunar Planet. Sci. 40. ID 1503) within the uncertainty. We found that in the phase angle range of 0° < α ? 25° for every 10° increase in phase angle Vesta’s visible slope (0.5–0.7 μm) increases 20%, Band I and Band II depths increase 2.35% and 1.5% respectively, and the BAR value increase 0.30. Phase angle spectral measurements of the eucrite Moama in the lab show a decrease in Band I and Band II depths and BAR from the lowest phase angle 13° to 30°, followed by possible small increases up to 90°, and then a dramatic drop between 90° and 120° phase angle. Temperature-induced spectral effects shift the Band I and II centers of the pyroxene bands to longer wavelengths with increasing temperature. We have derived new correction equations using a temperature series (80–400 K) of HED meteorite spectra that will enable interpretation of telescopic and spacecraft spectral data using laboratory calibrations at room temperature (300 K).     

The remarkable surface homogeneity of the Dawn mission target (1) Ceres

Dwarf-planet (1) Ceres is one of the two targets, along with (4) Vesta, that will be studied by the NASA Dawn spacecraft via imaging, visible and near-infrared spectroscopy, and gamma-ray and neutron spectroscopy. While Ceres’ visible and near-infrared disk-integrated spectra have been well characterized, little has been done about quantifying spectral variations over the surface. Any spectral variation would give us insights on the geographical variation of the composition and/or the surface age. The only work so far was that of Rivkin and Volquardsen ([2010], Icarus 206, 327) who reported rotationally-resolved spectroscopic (disk-integrated) observations in the 2.2–4.0 μm range; their observations showed evidence for a relatively uniform surface.Here, we report disk-resolved observations of Ceres with SINFONI (ESO VLT) in the 1.17–1.32 μm and 1.45–2.35 μm wavelength ranges. The observations were made under excellent seeing conditions (0.6″), allowing us to reach a spatial resolution of ～75 km on Ceres’ surface. We do not find any spectral variation above a 3% level, suggesting a homogeneous surface at our spatial resolution. Slight variations (about 2%) of the spectral slope are detected, geographically correlated with the albedo markings reported from the analysis of the HST and Keck disk-resolved images of Ceres (Li et al. [2006], Icarus 182, 143; Carry et al. [2008], Astron. Astrophys. 478, 235). Given the lack of constraints on the surface composition of Ceres, however, we cannot assert the causes of these variations.     

The surface energy balance at the Huygens landing site and the moist surface conditions on Titan

The Huygens Probe provided a wealth of data concerning the atmosphere of Titan. It also provided tantalizing evidence of a small amount of surface liquid. We have developed a detailed surface energy balance for the Probe landing site. We find that the daily averaged non-radiative fluxes at the surface are 0.7 W m^?2, much larger than the global average value predicted by McKay et al. (1991) of 0.037 W m^?2. Considering the moist surface, the methane and ethane detected by the Probe from the surface is consistent with a ternary liquid of ethane, methane, and nitrogen present on the surface with mole fractions of methane, ethane, and nitrogen of 0.44, 0.34, and 0.22, respectively, and a total mass load of ～0.05 kg m^?2. If this liquid is included in the surface energy balance, only a small fraction of the non-radiative energy is due to latent heat release (～10^?3 W m^?2). If the amount of atmospheric ethane is less than 0.6×10^?5, the surface liquid is most likely evaporating over timescales of 5 Titan days, and the moist surface is probably a remnant of a recent precipitation event. If the surface liquid mass loading is increased to 0.5 kg m^?2, then the liquid lifetime increases to ～56 Titan days. Our modeling results indicate a dew cycle is unlikely, given that even when the diurnal variation of liquid is in equilibrium, the diurnal mass variation is only 3% of the total liquid. If we assume a high atmospheric mixing ratio of ethane (>0.6×10^?5), the precipitation of liquid is large (38 cm/Titan year for an ethane mixing ratio of 2×10^?5). Such a flux is many orders of magnitude in excess of the photochemical production rate of ethane.     

Laboratory studies on the sputtering contribution to the sodium atmospheres of Mercury and the Moon

To ascertain the importance of sputtering by solar wind ions on the formation of a sodium exosphere around Mercury and the Moon, we have irradiated with 4 keV He ions, the Na bearing tectosilicates: albite, labradorite, and anorthoclase, as well as adsorbed Na layers deposited on albite and on olivine (a neosilicate that does not contain Na). Sodium at the surface and near surface (<40 Å) was quantified with X-ray photoelectron spectroscopy before and after each irradiation to determine the depletion cross section. We measured a cross section for sputtering of Na adsorbed on mineral surfaces, σ_s ≈ 1 × 10^?15 cm² atom^?1. In addition, mass spectrometric analyses of the sputtered flux show that a large fraction of the Na is sputtered as ions rather than as neutral atoms. These results have strong implications for modeling the sodium population within the mercurian and the lunar exospheres.     

The heat flow during the formation of ribbon terrains on Venus

Ribbons are regularly spaced, between 2 and 6 km, troughs that exist on venusian tesserae, which are mainly located in, and characterize to, venusian crustal plateaus. Independent of the geological or temporal relations with other features, regularly and similarly spaced ribbons on several tesserae strongly suggest a thermal control on the thickness of the deformed layer. This can be used to constraint the heat flow at the time of ribbon formation, which holds important implications for the viability of the hypotheses that address the origin and evolution of crustal plateaus. For a brittle–ductile transition ∼1–3 km deep (as proposed from ribbon spacing), realistic strain rates, and a present-day surface temperature of 740 K, the implied heat flow is very high, 130–780 mW m⁻². If Venus has experienced higher surface temperatures due to climate forcing by massive volcanism, then the heat flow could be greatly reduced. For surface temperatures of 850–900 K the heat flow is 190–560, 60–230 and 20–130 mW m⁻² for brittle–ductile transition depths of 1, 2 and 3 km, respectively. Heat flow values around 80–100 mW m⁻² are reasonable for venusian hotspots, based on terrestrial analogs, but hardly consistent with coldspot settings. High surface temperatures are also required to maintain the crustal solidus deeper than a few kilometers during the formation of ribbon terrains. For the obtained heat flows, a solidus deeper than ∼30 km (the likely mean value for the crustal thickness) is difficult to achieve. This suggests that a substantial proportion of the crust beneath crustal plateaus was emplaced subsequently to the time when ribbon terrains were formed. Alternatively, at that time a magma reservoir inside the crust could have existed.     

Spectral detectability of Ca- and Mg-sulfates in Martian bright soils in the 4–5 μm wavelength range

Chemical analyses of soil samples performed at different landing sites on Mars suggest the presence of sulfate minerals. These minerals are also thought to be present in the globally mixed Martian bright soils covering large areas of the planet. However, remote soil spectra have so far provided only tentative identification of sulfates regarding mineral types and abundances. This paper concentrates on the detectability of four Ca- and Mg-sulfates (anhydrite, gypsum, kieserite, hexahydrite) in the 4–5 μm range of Martian remote soil spectra. This spectral range is important for sulfate detection as most fine-grained sulfates exhibit significant absorption bands between 4 and 5 μm, independent of the texture of the host soils (e.g., loose powdered or cemented soils). Furthermore, this is the spectral range for which the Planetary Fourier Spectrometer (PFS) and Observatoire pour la Minéralogie, l’Eau, les Glaces, et l’Activité (OMEGA) instruments onboard ESA/Mars Express mission provide high spectral and spatial resolution data. Laboratory near- and mid-IR reflectance spectra of the pure sulfates and their mixtures with a terrestrial Martian soil analog were acquired. The results show that even the smallest amount of admixed sulfate (∼5 wt%) generates significant absorption features in the portion of the 4–5 μm range not covered by the saturated Martian atmospheric CO₂ absorption band between 4.2 and 4.4 μm. Model calculations of the influence of emitted surface radiation on the detectability of sulfate features show that the depth of the features decreases strongly with increasing surface temperature of an observed area resulting in the fact that all sulfates are spectrally hidden at surface temperatures around 270 K even at ∼14 or ∼25 wt% sulfate content in the soils. Sulfates become increasingly detectable depending on the sulfate content if the surface temperature is below 260 K. The outcome of this work helps to constrain the conditions needed for remote detection of sulfates within Martian bright soils in the 4–5 μm range.     

Interpretation of combined infrared,submillimeter, and millimeter thermal flux data obtained during the Rosetta fly-by of Asteroid (21) Lutetia

The European Space Agency’s Rosetta spacecraft is the first Solar System mission to include instrumentation capable of measuring planetary thermal fluxes at both near-IR (VIRTIS) and submillimeter–millimeter (smm–mm, MIRO) wavelengths. Its primary mission is a 1 year reconnaissance of Comet 67P/Churyumov–Gerasimenko beginning in 2014. During a 2010 close fly-by of Asteroid 21 Lutetia, the VIRTIS and MIRO instruments provided complementary data that have been analyzed to produce a consistent model of Lutetia’s surface layer thermal and electrical properties, including a physical model of self-heating. VIRTIS dayside measurements provided highly resolved 1 K accuracy surface temperatures that required a low thermal inertia, I < 30 J/(K m² s^0.5). MIRO smm and mm measurements of polar night thermal fluxes produced constraints on Lutetia’s subsurface thermal properties to depths comparable to the seasonal thermal wave, yielding a model of I < 20 J/(K m² s^0.5) in the upper few centimeters, increasing with depth in a manner very similar to that of Earth’s Moon. Subsequent MIRO-based model predictions of the dayside surface temperatures reveal negative offsets of ～5–30 K from the higher VIRTIS-measurements. By adding surface roughness in the form of 50% fractional coverage of hemispherical mini-craters to the MIRO-based thermal model, sufficient self-heating is produced to largely remove the offsets relative to the VIRTIS measurements and also reproduce the thermal limb brightening features (relative to a smooth surface model) seen by VIRTIS. The Lutetia physical property constraints provided by the VIRTIS and MIRO data sets demonstrate the unique diagnostic capabilities of combined infrared and submillimeter/millimeter thermal flux measurements.     

Enceladus: A hypothesis for bringing both heat and chemicals to the surface

The eruptive plumes and large heat flow (～15 GW) observed by Cassini in the South Polar Region of Enceladus may be expressions of hydrothermal activity inside Enceladus. We hypothesize that a subsurface ocean is the heat reservoir for thermal anomalies on the surface and the source of heat and chemicals necessary for the plumes. The ocean is believed to contain dissolved gases, mostly CO₂ and is found to be relatively warm (～0 °C). Regular tidal forces open cracks in the icy crust above the ocean. Ocean water fills these fissures. There, the conditions are met for the upward movement of water and the dissolved gases to exsolve and form bubbles, lowering the bulk density of the water column and making the pressure at its bottom less than that at the top of the ocean. This pressure difference drives ocean water into and up the conduits toward the surface. This transportation mechanism supports the thermal anomalies and delivers heat and chemicals to the chambers from which the plumes erupt. Water enters these chambers and there its bubbles pop and loft an aerosol mist into the ullage. The exiting plume gas entrains some of these small droplets. Thus, nonvolatile chemical species in ocean water can be present in the plume particles. A CO₂ equivalent-gas molar fraction of ～4 × 10^?4 for the ocean is sufficient to support the circulation. A source of heat is needed to keep the ocean warm at ～0 °C (about two degrees above its freezing point). The source of heat is unknown, but our hypothesis is not dependent on any particular mechanism for producing the heat.     

Comparisons of Cassini flybys of the Titan magnetospheric interaction with an MHD model: Evidence for organized behavior at high altitudes

Recent papers suggest the significant variability of conditions in Saturn’s magnetosphere at the orbit of Titan. Because of this variability, it was expected that models would generally have a difficult time regularly comparing to data from the Titan flybys. However, we find that in contrast to this expectation, it appears that there is underlying organization of the interaction features roughly above ～1800 km (1.7 Rt) altitude by the average external field due to Saturn’s dipole moment. In this study, we analyze Cassini’s plasma and magnetic field data collected at 9 Titan encounters during which the external field is close to the ideal southward direction and compare these observations to the results from a 2-fluid (1 ion, 1 electron) 7-species MHD model simulations obtained under noon SLT conditions. Our comparative analysis shows that under noon SLT conditions the Titan plasma interaction can be viewed in two layers: an outer layer between 6400 and 1800 km where interaction features observed in the magnetic field are in basic agreement with a purely southward external field interaction and an inner layer below 1800 km where the magnetic field measurements show strong variations and deviate from the model predictions. Thus the basic features inferred from the Voyager 1 flyby seem to be generally present above ～1800 km in spite of the ongoing external variations from SLT excursions, time variability and magnetospheric current systems as long as a significant southward external field component is present. At around ～1800 km kinetic effects (such as mass loading and heavy ion pickup) and below 1800 km ionospheric effects (such as drag of ionospheric plasma due to coupling with neutral winds and/or magnetic memory of Titan’s ionosphere) complicate what is observed.     

Winds and tides of Ligeia Mare,with application to the drift of the proposed time TiME (Titan Mare Explorer) capsule

We use two independent General Circulation Models (GCMs) to estimate surface winds at Titan’s Ligeia Mare (78° N, 250° W), motivated by a proposed mission to land a floating capsule in this ～500 km hydrocarbon sea. The models agree on the overall magnitude (～0.5–1 m/s) and seasonal variation (strongest in summer) of windspeeds, but details of seasonal and diurnal variation of windspeed and direction differ somewhat, with the role of surface exchanges being more significant than that of gravitational tides in the atmosphere. We also investigate the tidal dynamics in the sea using a numerical ocean dynamics model: assuming a rigid lithosphere, the tidal amplitude is up to ～0.8 m. Tidal currents are overall proportional to the reciprocal of depth—with an assumed central depth of 300 m, the characteristic tidal currents are ～1 cm/s, with notable motions being a slosh between Ligeia’s eastern and western lobes, and a clockwise flow pattern.We find that a capsule will drift at approximately one tenth of the windspeed, unless measures are adopted to augment the drag areas above or below the waterline. Thus motion of a floating capsule is dominated by the wind, and is likely to be several km per Earth day, a rate that will be readily measured from Earth by radio navigation methods. In some instances, the wind vector rotates diurnally such that the drift trajectory is epicyclic.     

Erosive flood events on the surface of Mars: Application to Mangala and Athabasca Valles

We address key factors involved in determining water flow conditions in outflow channels on Mars, including the temperature of the sub-surface water being released and the environmental conditions of low temperature, low atmospheric pressure, and low acceleration due to gravity. We suggest how some of the assumptions made in previous work may be improved. Our model considers the thermodynamic effects of simultaneous evaporation and freezing of water, and fluid dynamical processes including changes in flow rheology caused by assimilation of cold rock and ice eroded at the channel bed, and ice crystal growth due to water freezing. We model how far initially turbulent water could flow in a channel before it erodes and entrains enough material to become laminar, and subsequently ceases to erode the bed. An ice raft will begin to form on the flood while transition occurs between turbulent and laminar flow. Estimates are given for water transit times, ～17–19 h, initial water depths, 50–62 m, and average flow speeds, 5–12 m s^?1, in the Mangala and Athabasca Valles. We show that these two outflow channels, and by implication others like them, could plausibly have been formed in single water release events. Resulting mean erosion rates are approximately 0.7 mm s^?1, a factor of three greater than previous estimates based on combinations of estimates of flood duration and required water volumes. This is explained by the consideration of the effects of eroded ice and the physics of thermal erosion in the present study.     

Trapped Ar isotopes in meteorite ALH 84001 indicate Mars did not have a thick ancient atmosphere

Water is not currently stable in liquid form on the martian surface due to the present mean atmospheric pressure of ～7 mbar and mean global temperature of ～220 K. However, geomorphic features and hydrated mineral assemblages suggest that Mars’ climate was once warmer and liquid water flowed on the surface. These observations may indicate a substantially more massive atmosphere in the past, but there have been few observational constraints on paleoatmospheric pressures. Here we show how the ⁴⁰Ar/³⁶Ar ratios of trapped gases within martian meteorite ALH 84001 constrain paleoatmospheric pressure on Mars during the Noachian era [～4.56–3.8 billion years (Ga)]. Our model indicates that atmospheric pressures did not exceed ～1.5 bar during the first 400 million years (Ma) of the Noachian era, and were <400 mbar by 4.16 Ga. Such pressures of CO₂ are only sufficient to stabilize liquid water on Mars’ surface at low latitudes during seasonally warm periods. Other greenhouse gases like SO₂ and water vapor may have played an important role in intermittently stabilizing liquid water at higher latitudes following major volcanic eruptions or impact events.     

Low pressure and desiccation effects on methanogens: Implications for life on Mars

Conditions on the surface of Mars would appear to be too hostile for life as we know it. But the subsurface is another matter. If liquid water is present, even intermittently, life forms present would at least be protected from the lethal radiation bombarding the surface. However, life would have to contend with variations in pressure and possibly extended periods of desiccation. The research reported here involves both active metabolism (methanogenesis) at 400 and 50 mbar of pressure, pressures that would be found in the near subsurface of Mars, and survival following desiccation at both 1 bar (a pressure that would be found in the Martian subsurface) and 6 mbar (the lowest pressure at the surface and very near subsurface). The three methanogens tested for active metabolism, Methanothermobacter wolfeii, Methanosarcina barkeri and Methanobacterium formicicum, all demonstrated methane production at both 400 and 50 mbar on JSC Mars-1, a Mars soil simulant. Methane production at 50 mbar was much reduced compared to that at 400 mbar, most likely due to the greater stress at the lower pressure. In desiccation survival experiments, M. barkeri had survived 330 days of desiccation at 1 bar, while M. wolfeii and M. formicicum survived 180 and 120 days, respectively. Methanococcus maripaludis did not survive desiccation at all at 1 bar. At 6 mbar, M. wolfeii, M. barkeri and M. formicicum survived 120 days of desiccation while M. maripaludis survived 60 days. These results along with results from previous research would seem to indicate that there is no reason that methanogens could not inhabit the subsurface of Mars.     

Rotation period of Venus estimated from Venus Express VIRTIS images and Magellan altimetry

The 1.02 μm wavelength thermal emission of the nightside of Venus is strongly anti-correlated to the elevation of the surface. The VIRTIS instrument on Venus Express has mapped this emission and therefore gives evidence for the orientation of Venus between 2006 and 2008. The Magellan mission provided a global altimetry data set recorded between 1990 and 1992. Comparison of these two data sets reveals a deviation in longitude indicating that the rotation of the planet is not fully described by the orientation model recommended by the IAU. This deviation is sufficiently large to affect estimates of surface emissivity from infrared imaging. A revised period of rotation of Venus of 243.023 ± 0.002 d aligns the two data sets. This period of rotation agrees with pre-Magellan estimates but is significantly different from the commonly accepted value of 243.0185 ± 0.0001 d estimated from Magellan radar images. It is possible that this discrepancy stems from a length of day variation with the value of 243.023 ± 0.002 d representing the average of the rotation period over 16 years.     

Stardust-NExT,Deep Impact,and the accelerating spin of 9P/Tempel 1

The evolution of the spin rate of Comet 9P/Tempel 1 through two perihelion passages (in 2000 and 2005) is determined from 1922 Earth-based observations taken over a period of 13 year as part of a World-Wide observing campaign and from 2888 observations taken over a period of 50 days from the Deep Impact spacecraft. We determine the following sidereal spin rates (periods): 209.023 ± 0.025°/dy (41.335 ± 0.005 h) prior to the 2000 perihelion passage, 210.448 ± 0.016°/dy (41.055 ± 0.003 h) for the interval between the 2000 and 2005 perihelion passages, 211.856 ± 0.030°/dy (40.783 ± 0.006 h) from Deep Impact photometry just prior to the 2005 perihelion passage, and 211.625 ± 0.012°/dy (40.827 ± 0.002 h) in the interval 2006–2010 following the 2005 perihelion passage. The period decreased by 16.8 ± 0.3 min during the 2000 passage and by 13.7 ± 0.2 min during the 2005 passage suggesting a secular decrease in the net torque. The change in spin rate is asymmetric with respect to perihelion with the maximum net torque being applied on approach to perihelion. The Deep Impact data alone show that the spin rate was increasing at a rate of 0.024 ± 0.003°/dy/dy at JD2453530.60510 (i.e., 25.134 dy before impact), which provides independent confirmation of the change seen in the Earth-based observations.The rotational phase of the nucleus at times before and after each perihelion and at the Deep Impact encounter is estimated based on the Thomas et al. (Thomas et al. [2007]. Icarus 187, 4–15) pole and longitude system. The possibility of a 180° error in the rotational phase is assessed and found to be significant. Analytical and physical modeling of the behavior of the spin rate through of each perihelion is presented and used as a basis to predict the rotational state of the nucleus at the time of the nominal (i.e., prior to February 2010) Stardust-NExT encounter on 2011 February 14 at 20:42.We find that a net torque in the range of 0.3–2.5 × 10⁷ kg m² s^?2 acts on the nucleus during perihelion passage. The spin rate initially slows down on approach to perihelion and then passes through a minimum. It then accelerates rapidly as it passes through perihelion eventually reaching a maximum post-perihelion. It then decreases to a stable value as the nucleus moves away from the Sun. We find that the pole direction is unlikely to precess by more than ～1° per perihelion passage. The trend of the period with time and the fact that the modeled peak torque occurs before perihelion are in agreement with published accounts of trends in water production rate and suggests that widespread H₂O out-gassing from the surface is largely responsible for the observed spin-up.     

Locations of thin liquid water layers on present-day Mars

CRISM indicates the presence of water ice patches in Richardson crater, located on Mars’ southern polar region at the area of the seasonal ice cap. Numerical simulations suggest that the maximum daytime temperature of the ice at these locations is between 195 and 220 K during local spring. Previous studies suggest that at these temperatures liquid interfacial water could be present. Here, for the first time, we provide an example where the environmental conditions allow for the formation of such liquid films on present day Mars at the southern hemisphere. The upper bound estimated H₂O loss during the presence of these water ice patches is approximately 30 μm between Ls = 200 and 220, though it may be as low as 0.1 μm depending on the ambient water vapor. The upper bound value is larger than the expected condensation thickness in autumn; however, it may still be realistic due to CO₂ gas jet generated deposition and possible subsequent accumulation on mineral grains. The presence of this interfacial water may have impact on local chemical processes along with astrobiological importance.     

The planetary fourier spectrometer (PFS) onboard the European Venus Express mission

The planetary fourier spectrometer (PFS) for the Venus Express mission is an infrared spectrometer optimized for atmospheric studies. This instrument has a short wavelength (SW) channel that covers the spectral range from 1700 to 11400 cm⁻¹ (0.9–5.5 μm) and a long wavelength (LW) channel that covers 250–1700 cm⁻¹ (5.5–45 μm). Both channels have a uniform spectral resolution of 1.3 cm⁻¹. The instrument field of view FOV is about 1.6 ° (FWHM) for the short wavelength channel and 2.8 ° for the LW channel which corresponds to a spatial resolution of 7 and 12 km when Venus is observed from an altitude of 250 km. PFS can provide unique data necessary to improve our knowledge not only of the atmospheric properties but also surface properties (temperature) and the surface-atmosphere interaction (volcanic activity).PFS works primarily around the pericentre of the orbit, only occasionally observing Venus from larger distances. Each measurements takes 4.5 s, with a repetition time of 11.5 s. By working roughly 1.5 h around pericentre, a total of 460 measurements per orbit will be acquired plus 60 for calibrations. PFS is able to take measurements at all local times, enabling the retrieval of atmospheric vertical temperature profiles on both the day and the night side.The PFS measures a host of atmospheric and surface phenomena on Venus. These include the:(1) thermal surface flux at several wavelengths near 1 μm, with concurrent constraints on surface temperature and emissivity (indicative of composition); (2) the abundances of several highly-diagnostic trace molecular species; (3) atmospheric temperatures from 55 to 100 km altitude; (4) cloud opacities and cloud-tracked winds in the lower-level cloud layers near 50-km altitudes; (5) cloud top pressures of the uppermost haze/cloud region near 70–80 km altitude; and (6) oxygen airglow near the 100 km level. All of these will be observed repeatedly during the 500-day nominal mission of Venus Express to yield an increased understanding of meteorological, dynamical, photochemical, and thermo-chemical processes in the Venus atmosphere. Additionally, PFS will search for and characterize current volcanic activity through spatial and temporal anomalies in both the surface thermal flux and the abundances of volcanic trace species in the lower atmosphere.Measurement of the 15 μm CO₂ band is very important. Its profile gives, by means of a complex temperature profile retrieval technique, the vertical pressure-temperature relation, basis of the global atmospheric study.PFS is made of four modules called O, E, P and S being, respectively, the interferometer and proximity electronics, the digital control unit, the power supply and the pointing device.     

The XMM-Newton view of the yellow hypergiant IRC + 10420 and its surroundings

Among evolved massive stars likely in transition to the Wolf–Rayet phase, IRC + 10420 is probably one of the most enigmatic. It belongs to the category of yellow hypergiants and it is characterized by quite high mass loss episodes. Even though IRC + 10420 benefited of many observations in several wavelength domains, it has never been a target for an X-ray observatory. We report here on the very first dedicated observation of IRC + 10420 in X-rays, using the XMM-Newton satellite. Even though the target is not detected, we derive X-ray flux upper limits of the order of 1–3 × 10⁻¹⁴ erg cm⁻² s⁻¹ (between 0.3 and 10.0 keV), and we discuss the case of IRC + 10420 in the framework of emission models likely to be adequate for such an object. Using the Optical/UV Monitor on board XMM-Newton, we present the very first upper limits of the flux density of IRC + 10420 in the UV domain (between 1800 and 2250 Å and between 2050 and 2450 Å). Finally, we also report on the detection in this field of 10 X-ray and 7 UV point sources, and we briefly discuss their properties and potential counterparts at longer wavelengths.